Structure and evolution of single WR stars

New computations of massive stars follow the evolution up to advanced stages and include:

-A large and flexible nuclear network consisting of 174 nuclear species that are linked by 1742 nuclear reactions.

-Semiconvection, overshooting and mass loss.

-Modern rates for both strong and weak interaction processes as well as the latest rates for the neutrino processes.

-Improved grid distribution and a large number of grid points.

The nuclear network and the diffusion equation are solved for each time step during the whole evolution. In this way the accuracy of nuclear yields and chemical abundances are mainly limited by uncertainties in the diffusion coefficient found from the convection theories. Several instability mechanisms may affect the mass loss rates of massive stars and thereby the structure and abundances of WR stars. Due to heavy mass loss at the LBV and WR stages, the masses at the pre-SN stage may be less than 5M _⊙. Yields and abundances throughout the stars are discussed together with the amount of all elements expelled.     

Cosmic Ray Anisotropies Around the Forbush Decrease of April 11, 2001

During the Forbush decrease of April 11, 2001, the Ooty (11.4^∘N, 76.7^∘E, 2200 m altitude, South India) high energy muon detectors showed considerable anisotropies. Some anisotropies and asymmetries indicated that the Earth did not pass the middle of the interplanetary structure (blob), and passed its northern part. Some anisotropies which occurred when the Earth was outside the blob (notably before the Forbush decrease) could be the precursory increases due to reflection from the shock fronts, but some others could not be understood as these appeared in a direction away from the blob.     

Sunspots with the Strongest Magnetic Fields

The strongest observed solar magnetic fields are found in sunspot umbrae and associated light bridges. We investigate systematic measurements of approximately 32 000 sunspot groups observed from 1917 through 2004 using data from Mt. Wilson, Potsdam, Rome and Crimea observatories. Isolated observations from other observatories are also included. Corrections to Mt. Wilson measurements are required and applied. We found 55 groups (0.2%) with at least one sunspot with one magnetic field measurement of at least 4000 G including five measurements of at least 5000 G and one spot with a record field of 6100 G. Although typical strong-field spots are large and show complex structure in white light, others are simple in form. Sometimes the strongest fields are in light bridges that separate opposite polarity umbras. The distribution of strongest measured fields above 3 kG appears to be continuous, following a steep power law with exponent about −9.5. The observed upper limit of 5 – 6 kG is consistent with the idea that an umbral field has a more or less coherent structure down to some depth and then fragments. We find that odd-numbered sunspot cycles usually contain about 30% more total sunspot groups but 60% fewer >3 kG spots than preceding even-numbered cycles.     

Comparison of the Variations of CMEs and ICMEs with those of other Solar and Interplanetary Parameters During Solar Cycle 23

This paper examines the variations of coronal mass ejections (CMEs) and interplanetary CMEs (ICMEs) during solar cycle 23 and compares these with those of several other indices. During cycle 23, solar and interplanetary parameters had an increase from 1996 (sunspot minimum) to ∼2000, but the interval 1998–2002 had short-term fluctuations. Sunspot numbers had peaks in 1998, 1999, 2000 (largest), 2001 (second largest), and 2002. Other solar indices had matching peaks, but the peak in 2000 was larger than the peak in 2001 only for a few indices, and smaller or equal for other solar indices. The solar open magnetic flux had very different characteristics for different solar latitudes. The high solar latitudes (45^∘–90^∘) in both N and S hemispheres had flux evolutions anti-parallel to sunspot activity. Fluxes in low solar latitudes (0^∘–45^∘) evolved roughly parallel to sunspot activity, but the finer structures (peaks etc. during sunspot maximum years) did not match with sunspot peaks. Also, the low latitude fluxes had considerable N–S asymmetry. For CMEs and ICMEs, there were increases similar to sunspots during 1996–2000, and during 2000–2002, there was good matching of peaks. But the peaks in 2000 and 2001 for CMEs and ICMEs had similar sizes, in contrast to the 2000 peak being greater than the 2001 peak for sunspots. Whereas ICMEs started decreasing from 2001 onwards, CMEs continued to remain high in 2002, probably due to extra contribution from high-latitude prominences, which had no equivalent interplanetary ICMEs or shocks. Cosmic ray intensity had features matching with those of sunspots during 2000–2001, with the 2000 peak (on a reverse scale, actually a cosmic ray decrease or trough) larger than the 2001 peak. However, cosmic ray decreases started with a delay and ended with a delay with respect to sunspot activity.     

A comparison of type II shock speeds

Type II radio bursts are produced by material moving outwards in the solar atmosphere. Their drift in frequency allows the calculation of the radial speed with which the shock is moving- very basic information in assessing the likelihood that the shock will reach the Earth and its time of arrival. This paper compares the shock speeds derived from radio bursts observed by the Swept Frequency Interferometric Radiometer (SFIR) equipment at the US Air Force Radio Solar Telescope Network (RSTN) of observatories with those measured with the Culgoora radiospectrograph operated by IPS Radio and Space Services. The SFIR shock speeds are found to be 1.5–3.0 times larger than the Culgoora values which are consistent with earlier results. This difference appears to originate from the incorrect interpretation of events as a result of the smaller frequency range of the SFIR equipment.     

Whistler observations of the quiet time plasmasphere-ionosphere coupling fluxes at low latitude

Observations of whistlers during quiet times made at low-latitude ground station Nainital (geomag. lat. 19 1 N) are used to deduce plasmasphere-ionosphere coupling fluxes. The whistler data from 3 magnetically quiet days are presented that show a smooth decrease in dispersion with time. This decrease in dispersion is interpreted in terms of a corresponding decrease in electron content of tubes of ionization. The electron densities, electron tube contents (10¹⁶ el/m²-tube) and coupling fluxes (10 el m^–1 s^–2) are computed by means of an accurate curve fitting method developed by Tarcsai (1975) and are in good agreement with the results reported by other workers.     

First results from the Gauribidanur Radio heliograph

Observations of the Sun at two frequencies (51 and 77 MHz) using the East-West arm of the Gauribidanur Radio heliograph are presented.     

Center-to-limb variation of low-frequency Ca II line brightness oscillations in the solar chromosphere

The goal of this paper is a detailed statistical analysis of the low-frequency Ca II line intensity oscillations containing information about the dynamics of the lower and middle chromosphere. A pixel-by-pixel analysis of the observed parameters has been performed. The following results have been obtained. (1) The low-frequency chromospheric oscillations (periods >400 s) are seen much more frequently in networks than in chromospheric network cells. (2) The relative fraction of the low-frequency chromospheric intensity oscillations increases with height. (3) The occurrence distribution of intensity oscillations as a function of the frequency is subdivided at least into two types. (4) In contrast to the low-frequency photospheric oscillations, the phase differences between the Ca II K and 849.8 nm line intensity oscillations do not give grounds to identify the low-frequency chromospheric oscillations with internal gravity waves. (5) The spectral composition of the oscillations in the network chromosphere resembles that expected in magnetic flux tubes in the nonlinear regime of conversion of transverse MHD waves at lower levels of the atmosphere into longitudinal MHD waves in its upper layer.